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The following handle holds various files of this Leiden University dissertation:
http://hdl.handle.net/1887/80839
Author: Haffert, S.Y.
Almost three decades ago our perception of the universe changed dras-tically. The first planet around a star other than our own Sun had been found (Wolszczan & Frail, 1992). This system was and still is rather unique because the two planets are orbiting a pulsar, the remnant of a star after it has created a supernova. This fact puzzled astronomers because there were no known methods at that time for planets to survive a supernova. Another possibility was that they formed from the left-over debris (Rasio et al., 1992; Tavani & Brookshaw, 1992). The second shock came when the first exoplanet orbiting a solar-like star was discovered just three years later (Mayor & Queloz, 1995). The planet, 51 Pegasi b, is in a very short-period orbit of only 4.23 days and roughly half the mass of Jupiter. It was very surprising to find a planet comparable to Jupiter orbiting their host star much closer than that Mercury is orbiting around the Sun. More Jupiter-like planets on close-in orbits followed soon (Butler et al., 1997; Marcy & Butler, 1996). This class of gas-giant planets was quickly termed ’hot Jupiters’ because the close proximity to their host star leads to high equilibrium temperatures.
In the years after these first few discoveries the field of exoplanet re-search quickly expanded. Many observing techniques and instruments were developed, leading to an explosive growth in the number of discovered plan-ets, which can be seen in Figure 1.1. Most exoplanets to date have been found by the Kepler mission, which added almost 2500 planets. The Kepler mission used the transit method where stars are closely monitored to search for periodic dimmings when the planet moves in front of the star (Borucki et al., 2010; Henry et al., 2000). Kepler revealed that there are many ex-otic planets and planetary systems. A surprising find was the detection of many super-Earths and sub-Neptunes with masses of a few times that of the Earth (Petigura et al., 2013a,b). These types of planets are the most ubiquitous in the Milky Way even though our own Solar system does not have any of them (Petigura et al., 2013a,b).
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This variety in exoplanets and planetary systems is challenging the the-ories of planet formation because the entire range of observed planetary-system architectures must be explained. The initial conditions for planet formation are set by the formation of the host star. Therefore the formation of planets cannot be understood independently from star formation. Stars are formed from clouds of molecular gas in the interstellar medium. Small overdensities in these large, cold clouds can create gravitational instabil-ities thath lead to the local collapse of the gas clouds into proto-stellar cores (McKee & Ostriker, 2007). If the collapsing gas has some angular momentum, it will flatten out the collapsing cloud and form a disk with the proto-star at its center. The surrounding dust and gas will gather into the circumstellar disk, which is thought to be the birth place of planets and is therefore also called a protoplanetary disk (Armitage & Belmonte, 2018). There are several proposed mechanism through which planets can form, and they broadly fall into one of the following three categories:
1. The planet forms through core accretion where small dust particles slowly coagulate into a proto-planetary core (Pollack et al., 1996). As the core grows, its gravity also grows, and it will attract more dust. When the proto-planet is massive enough it will start to rapidly ac-crete the gas and dust in its surrounding, thereby clearing out a path in the circumstellar disk through runaway accretion. This process stops when the proto-star becomes luminous enough to clear the disk through radiative pressure.
2. There are several mechanism through which the protoplanetary disk can become unstable and fragment into self-gravitating clumps. The most common method proposed for this are gravitational instabilities Boss (1997) that are created if the disk is very massive. But recent ALMA observations have revealed that massive disks are not very common, and this makes the gravitational instability process possi-bly a very rare event (Andrews et al., 2013; Pascucci et al., 2016). In the last few years it has been argued that magneto-rotational insta-bilities (MRI) may also cause disk fragmentation that leads to planet formation (Chiang & Youdin, 2010).
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It is possible that all three processes play a part in the formation of planets. One of the challenges will be to determine which process dominates the formation process for which class of planets. It has been suggested that gas giants onwide orbits like those in the HR8799 system have been formed through gravitational instability (Nero & Bjorkman, 2009). However there is also contradicting evidence that their masses and separation do not fulfill several of the criteria for the formation through such instabilities (Bowler et al., 2015; Rameau et al., 2013; Vorobyov, 2013).
The interaction between the planet and the protoplanetary disk is thought to be quite complex (Kley, 2017). The planet can change its orbital dis-tance, either moving in or out, due to planet-disk interaction. The more massive planets are able to sweep up a major part of the disk material in their orbit and carve a deep gap in the disk. The depletion of the dust and gas in the disk changes the pressure gradient and forces the planet to mi-grate; this migration scenario is called type-I migration (Kley, 2017; Nelson et al., 2000). Planets of a few Earth masses follow a different migration scenario called type II (Nelson et al., 2000) where only a small shallow gap is created that is not completely cleared of dust and gas. The main differ-ence between the different types is the amount of matter that is accreted, and that determines whether the planet-disk interaction is linear (type II) or non-linear (type I). The case for multiple planets is more complicated since the planets will also influence each other, which is classified as type-III migration. In the past decades complex hydro-dynamical simulations have been conducted to understand the behaviour of migrating planets, leading to the development of semi-analytical relations between the migration rate, disk parameters and planet parameters (Dodson-Robinson & Salyk, 2011; Kley, 2017).
most of the physical mechanisms in these population codes (Morbidelli & Raymond, 2016).
Direct imaging plays an important role to overcome these observational limitations. By spatially resolving the disk and the embedded planets, we can witness their interaction. Another added benefit is the enhanced in-trinsic contrast between the star and the planets. For old systems such as our own solar system, the best signal we could use to detect Earth or Jupiter from a distance is through reflected light. The intensity ratio be-tween the Sun and the reflected light of Earth and Jupiter are 10−10 and 10−9, respectively (Traub & Oppenheimer, 2010). This is a huge contrast to overcome. But during the first stages of planet formation, the planets are still very hot. This increases the intrinsic contrast in the Near-Infrared to 10−5− 10−6 (Burrows et al., 2004) making the detection of such exo-planets orders of magnitude easier. This shows that direct imaging is the prime technique to observe young planetary systems and their planet-disk interactions.
1.1
The direct imaging challenge
1.1.1 The Earth atmosphere
Direct imaging of exoplanets is a challenging task because a high contrast needs to be reached at very close angular separations. If we place our solar system at 100 parsec, the resolving power necessary to separate Earth from the Sun would need to be better than 10 milliarcseconds (mas), but even if we could resolve Earth, the contrast between the Earth and the Sun of about 10−10 will make Earth close to impossible to observe. For Jupiter it becomes slightly easier with a separation of 55 mas and a contrast of 10−8− 10−9. To resolve Earth and Jupiter at this distance we would
need to use large telescopes of at least 30 meters in diameter, under the assumption that we will be able to solve the contrast-ratio problem. This angular resolving power will become available in the next decade with the construction of the upcoming extremely large telescopes; the Extremely Large Telescope (ELT) spearheaded by ESO, the Thirty Meter Telescope (TMT) and the Giant Magellan Telescope (GMT). But until those are build, we will have to use the current 8 and 10-meter class telescopes that are limited to about 26 mas angular resolution at 1 µm by diffraction,
The direct imaging challenge 6
Here ∆θ is the angular resolving power, λ the wavelength that is used for imaging and D the telescope diameter. While the current generation of telescopes like the Very Large Telescope (VLT) of ESO, with an 8.2-meter diameter, would be able to resolve Jupiter at 100 pc, we have not been able to do this. For ground-based telescopes there are two challenges to overcome. The first being turbulence in the Earth’s atmosphere, and the second is the intrinsic contrast between the planet and its host star. When light propagates from a star towards the Earth, it becomes a smooth plane wave due to the large distance between us and the star. It travels over sev-eral years to tens or hundreds of year,s and when it finally reaches Earth the light has to travel through the atmosphere to enter our telescopes. Dur-ing the last tenth of a milliseconds of its journey the light wave loses its flatness because of turbulence in the atmosphere (Fried, 1966). This tur-bulence will create wavefront aberrations that degrade the resolving power of the telescope. The amount of wavefront aberration depends on the tur-bulence strength that is parametrised by the Fried parameter r0 (Fried,
1966). The Fried parameter is the characteristic spatial scale of the per-turbed wavefront where the wavefront changes by less than one radian. The resolution limit of the telescope is set by this characteristic scale instead of the telescope diameter. In median weather conditions the Fried parameter is roughly 20 to 30 cm at 1 µm for good observing sites such as Paranal, La Palma or Mauna Kea. The resolution that the VLT achieves during these condition is about 1 arcsecond, almost 40 times larger than the diffraction limit! This can be seen in Figure 1.2.
1.1.2 Adaptive optics
Diffraction-limited Short exposure Long exposure
The direct imaging challenge 8
Figure 1.3: These two figures show different types of adaptive optics. The left uses the light from the astrophysical target itself to do wavefront sensing, while the right scheme uses an artificial light source created by a powerful laser that is reflected by the upper atmosphere. Both methods drive a single deformable mirror to correct for the wavefront aberrations. Image credit: ESO.
modulations on the detector. The standard AO system as drawn in Figure 1.3 uses a WFS to measure wavefront deviations and feeds those back to the DM to create a closed-loop feedback system. The AO system needs to operate at several hunderd Hz to several thousand Hz because of the time scale over which the atmosphere changes (Greenwood, 1977). The coher-ence time of the atmosphere τ0 is roughly r0, the Fried parameter, divided
by the wind speed v (Greenwood, 1977). This leads to a coherence time on the order of 1 ms to 10 ms, which is why AO systems need to do the corrections in real time.
relatively easy to implement and has the potential to provide the highest possible on-axis correction. Many astronomers used these early AO instru-ments for direct imaging while they had not been specifically developed for exoplanet science (Chauvin et al., 2005). The potential of AO instruments for direct imaging was proven by the detection of multiple planets around HR8799 (Marois et al., 2008, 2010). This system came as a surprise because most planets found until then were much closer to their host star, making HR8799 a still unique planetary system.
SCAO has worked very well for the purpose of improving the image quality but it is limited to bright targets because the light from the star itself is used to measure the wavefront errors created by the atmosphere. In the past two decades a large amount of work has been done to make AO-corrected images accessible for fainter targets. Instead of using the light from the astrophysical object an artificial light source is generated with a powerful laser high up in the atmosphere (Foy & Labeyrie, 1985; Fugate et al., 1991). For large telescopes a sodium laser is used to excite atoms in the sodium layer of Earth’s atmosphere (Bonaccini Calia et al., 2010). The excited atoms will become an articial light beacon that can be used to measure the atmospheric turbulence. Due to the brightness of the laser it is not possible to bring the laser close to the astrophysical source, it needs to be pointed slightly away from the target. The atmospheric volume that is probed by this laser is slightly different than the volume that the star passes through. This led to the development of Laser Tomography Adaptive Op-tics (LTAO) where multiple laser guide-stars are placed around the target of interest (Hubin et al., 2005; Tallon & Foy, 1990). The measurements from the different lasers are then combined to create the best estimate of the on-axis wavefront errors. ESO applied this in the Adaptive Optics Fa-cility (AOF) for the VLT that saw first light in 2015 (Madec et al., 2018). It has since then produced spectacular images, see for example Fig 1.4.
1.1.3 High-contrast imaging
The direct imaging challenge 10
Figure 1.4: Observations of Neptune with the new Narrow Field Mode of MUSE with AO correction provided by the LTAO system. The LTAO shows almost diffraction-limited performance. Image credits to ESO/P. Weilbacher (AIP).
Optics (XAO). With the current generation of high-contrast imagers (HCI) we can reach diffraction-limited performance in the near-infrared. But this is not enough to find faint planets as the planet is still much fainter than the Airy rings of the stellar diffraction pattern. With the high quality of the PSFs of SPHERE and GPI they can also use advanced coronagraphs to remove the diffraction effects of the star.
masks have been developed that theoretically can remove all starlight if the input wavefront has no aberrations (Foo et al., 2005; Guyon, 2003; Rouan et al., 2000; Soummer, 2005). Another class of coronagraphs called pupil plane coronagraphs place masks in the pupil of the telescope to modify the shape of the PSF. By manipulating the amplitude or phase in the pupil, the electric field in the focal plane can be made to destructively interfere. With this technique dark holes can be created where we can search for planets, and because the optics are in the pupil, they are insensitive to vibrations. There are currently two flavours of pupil-plane coronagraphs, the Shaped Pupil (SP) coronagraph that uses amplitude masks (Kasdin et al., 2003; Soummer et al., 2003) and the Apodizing Phase Plate (APP) coronagraph that uses phase plates (Codona et al., 2006; Otten et al., 2017; Snik et al., 2012).
For Lyot-style coronagraphs the PSF needs to be perfectly aligned with the focal plane mask to cancel the starlight, but due to vibrations and small drifts the star will not be perfectly aligned with the mask. This deterio-rates the performance of the coronagraph (Ruane et al., 2017). Pupil-plane coronagraphs are less sensitive to theses issues because the optical elements are in the pupil. Next to vibrations all other wavefront errors will also de-grade the performance of the coronagraph (Aime & Soummer, 2004). There are still residual wavefront errors even though an AO system is used. The residual wavefront errors have two sources, the first being residual wave-front errors from the atmosphere that are not correctable or not completely removed. The second is due to a difference in the optical path between the coronagraphic optics and the wavefront-sensor optics. Because these in-struments have different optics, they will see a slightly different wavefront error causing differential wavefront errors between the two systems. These wavefront errors are called Non-Common Path Aberrations (NCPAs). A lot of current research is focused on mitigating these NCPAs (Jovanovic et al., 2018). Both the NCPAs and the residual turbulence causes speckles that can look like planets. Image-processing algorithms are used to further remove these speckles.
1.1.4 Post-processing
The direct imaging challenge 12
object and then subtract its PSF from the science target. Because this reference object is used to measure the PSF, it should not include any circumstellar material or companions. This technique is called Reference Differential Imaging (RDI) and was one of the first HCI techniques and was able to reveal the circumstellar disk around Beta Pictoris (Smith & Terrile, 1984). RDI has also been very successfully applied to Hubble Space Telescope (HST) data because HST has a very stable PSF (Schneider & Silverstone, 2003). If it is not possible to use a reference target, either due to unavailability or because the speckle pattern is not repeatable for different targets, a PSF model has to be built from the data itself. To create the reference PSF in this case, one needs to make use of a difference between the star and the planet.
Figure 1.5: An observation of the HR8799 system taken with the LBT (Maire et al., 2015) and post-processed with Angular Differential Imag-ing. The four planets orbiting the star are clearly resolved. Searching for planets closer in is difficult due to slowly changing speckles that limit how close in we can search. The speckle noise can be seen at the edge of the coronagraphic mask were the intensity quickly changes from white to black.
where limited improvement is achieved close to the star. Diversities that are based on the intrinsic properties of the observed system that are time invariant would be more robust against these varying speckles.
1.1.5 The powers of ten in exoplanet spectroscopy
The direct imaging challenge 14
signposts of a planet in formation (Aoyama et al., 2018; Marleau et al., 2017; Zhu, 2015), which occurs when gas is deposited onto the planet at high velocity. As the gas collides with the planet, it creates a strong shock front, which heats up the local gas to high temperatures (T>10000 K). This process generates a large amount of Hα emission, which decreases the con-trast between the star and planet by several orders of magnitude, thereby making it easier to detect. The difference between two narrowband filters with one covering Hα and the other in the nearby continuum can be used to subtract out the star (Close et al., 2014).
A higher-resolution version of this is Spectral Differential Imaging, (hav-ing the same abbreviation as Simultaneous Differential Imag(hav-ing). With SDI the PSF is measured at many wavelengths, usually with a low-resolution integral-field spectrograph at a resolving power of R = 50 − 100 over a large bandwidth. Due to the properties of diffraction the PSF and its speckles scale radially with wavelength while the planet stays at a fixed position (Sparks & Ford, 2002; Thatte et al., 2007). Rescaling the data to a ref-erence wavelength will overlay the speckles while smearing out the planet. Taking a median as is done with ADI will create a PSF model that can be used for subtraction. Some planet signal is also subtracted by this proce-dure; the amount of planet subtraction depends on the observed bandwidth and the angular distance of the planet. SDI has the advantage that it can remove the starlight and at the same time characterize the planet at low resolving power. This is very powerful because it provides a spectrum of the planet. Usually both SDI techniques are combined with ADI into sADI to make use of both diversities at the same time. The combined technique of sADI has allowed us to reach the deepest contrasts ever observed (Vigan et al., 2015).
16100 16200 16300 16400 16500 16600 16700 16800 16900 Wavelength (Å) 0.5 0.0 0.5 1.0 1.5 2.0 2.5 No rma lize d inte nsi ty R=100000 R=10000 R=1000 R=100 Planet spectrum Star spectrum
Figure 1.6: The spectrum of a solar-like star modelled with a 6000 K PHOENIX model and a spectrum of a giant planet modelled by a 1200 K BTSettl model. The resolving power changes by one order of magnitude between the different spectra, going from R=100000 to R=100. The spectra are shifted for ease of viewing. As the spectral resolving power decreases, it becomes more difficult to discriminate the planetary spectrum from the stellar spectrum.
advantage of this technique is that it is not limited by speckle noise, which hampers the other post-processing techniques.
Thesis outline 16
1.2
Thesis outline
The goal of this thesis is to explore the potential of high-resolution integral-field spectroscopy behind a high-contrast imaging instrument for the de-tection and characterization of exoplanets. The work presented in this thesis can be divided into three parts, the first one focused on coupling a high-contrast imager with a high-resolution spectrograph (R≈100000). The second part shows the scientific gain of integral-field spectroscopy in the visible for high-contrast imaging. And the last part is about a novel way to do spectroscopy with applications for astronomy and Earth observations.
Chapter 2 and 3: The Leiden EXoplanet Instrument(LEXI) These two chapters present the design, development and on-sky results of the Leiden EXoplanet Instrument (LEXI) a bench-mounted visitor in-strument for the 4.2m William Herschel Telescope at La Palma. LEXI was built as a test bed for high-contrast imaging and integral-field spec-troscopy. Several different approaches to AO-fed spectroscopy have been tested with LEXI. Our results show that XAO systems are well suited for single-mode fiber spectroscopy. LEXI has also been used to test several wavefront sensing concepts such as the generalised Optical Differentiation Wavefront Sensor (g-ODWFS) (Haffert (2016), Haffert et. al. in prep.), the Coronagraphic Modal Wavefront Sensor (Wilby et al., 2016, 2017) and more recently the Three Wave Shearing Interferometer (TWSI) (Por et al. in prep.).
Chapter 4 and 5: SCAR
These two chapters present the Single-mode Complex Amplitude Re-finer (SCAR) coronagraph. SCAR is a promising new coronagraph that makes use of the mode-filtering capabilities of single-mode fibers. This allows us to design and create coronagraphs with higher planet through-put that can search closer to the star. In chapter 5 we present the concept, designs and performance estimates where we show that SCAR enables coro-nagraphs with inner-working angles close to the diffraction limit. In chapter 6 we experimentally demonstrate the nulling capabilities of SCAR for two differently designs in the lab where we reached a 10−4 contrast at 1 λ/D.
Chapter 6: Imaging a forming multi-planet system
medium-resolution integral-field unit that spans the wavelength range from 0.465µm to 0.93µm at an average resolving power of R = λ/∆λ = 3000. The in-strument is fed by the LTAO system on UT4 of the VLTs and can reach a spatial resolution of roughly 60 milliarcseconds in good seeing conditions. The combination of the spectral resolving power and the AO performance made it possible to detect two accreting proto-planets in the transition disk around PDS 70. Our observations show that adaptive-optics-assisted, medium-resolution, integral-field spectroscopy with MUSE targeting ac-cretion signatures is a powerful way to trace ongoing planet formation in transitional disks at different stages of their evolution. This was also the first time that a planet has been discovered with an LTAO system, which is very interesting as LTAO can reach a better performance on fainter targets than comparable SCAO systems.
Chapter 7: Novel spectroscopic instrumentation
This chapter presents a novel spectrograph concept based on Volume Bragg Gratings (VBG) that is able to achieve high spectral resolution over a large wavelength range for a large field of view without the need for very large detectors. This is achieved by creating specialized spectral filters with highly multiplexed VBGs (HMBG) that are sensitive to a molecular species of choice. The HMBG condenses the full spectrum into a small, multiplexed spectrum with the size of a single spectral line thereby enabling a large re-duction of the required detector real estate per spatial pixel. The chapter presents the concept and a few case studies.
1.3
Outlook
Outlook 18
optically (Chapter 7), and therefore we can reduce the number of required detector pixels per spatial point. With the multiplexed Bragg gratings we can apply the same technique to much larger fields of view and bypass the field-of-view limitation of high-resolution integral-field spectroscopy.
Medium to high-resolution integral-field spectroscopy is likely to be the ideal observing technique to search for accretion signatures from proto-planets. The current standard is to search for Hα emission with Simul-taneous Differential Imaging, effectively resulting in resolving powers on the order of 10 − 100. Signatures such as Hα are intrinsically narrowband, therefore increasing the spectral resolving power of our observations in-creases the signal-to-noise ratio as long as the line is not resolved. Adding the capability to observe these signatures at much higher resolving power R = 5000 − 10000 will increase the signal-to-noise ratio by a factor 10-100 almost for free. MUSE at the VLT does have the capability of integral-field spectroscopy but it was not designed for high-contrast imaging, and therefore lacks the capability for starlight suppression. Development of high-resolution integral-field units for extreme adaptive optics systems will allow us to take the next step in the search and characterization of proto-planets, where we will be able to not only find such planets more efficiently but also can study the process of accretion in detail.
Currently MUSE provides an exciting opportunity to study the time variability of accretion signals from short to long timescales. Such obser-vations will set strong constraints on planet growth and evolution during the earlier stages. In addition due to the unique broad spectral coverage of MUSE, we can observe other accretion tracers such as Hβ at 4861˚A, OI at 8446˚A, and the CaII triplet at 8498˚A, 8542˚A, and 8662˚A. Together with Hα, the detection of any these tracers will put constraints on the temper-ature, density and shock velocity at the interface between the planet and the accretion flow.
This work at medium resolution lays down the foundation for visible-light high-resolution integral-field units and high-contrast imaging for the detection of reflected light from cold and old exoplanets, like Earth, and biosignatures such as the O2 band with the Extremely Large Telescopes
by using one of the ELTs. ELTs come with two advantages, the first being the larger collecting area, and the second is the increased spatial resolu-tion. With an ELT the detection of Proxima Centauri b can obtained in a single night instead of the hundred nights of VLT time (Snellen et al., 2015). With the addition of high-resolution integral-field units to extreme adaptive optics systems at ELTs, we will start to study older, potentially habitable planets, and thus address humanity’s ultimate question: Are we alone?
Bibliography
Aime, C., & Soummer, R. 2004, ApJ, 612, L85
Andrews, S. M., Rosenfeld, K. A., Kraus, A. L., & Wilner, D. J. 2013, ApJ, 771, 129
Aoyama, Y., Ikoma, M., & Tanigawa, T. 2018, ApJ, 866, 84
Armitage, P. J., & Belmonte, J. A. 2018, A Brief Overview of Planet Formation (Cham: Springer International Publishing), 2185–2203. https://doi.org/10. 1007/978-3-319-55333-7_135
Babcock, H. W. 1953, PASP, 65, 229
Baraffe, I., Chabrier, G., Barman, T. S., Allard, F., & Hauschildt, P. H. 2003, A&A, 402, 701
Barman, T. S., Konopacky, Q. M., Macintosh, B., & Marois, C. 2015, ApJ, 804, 61
Benz, W., Ida, S., Alibert, Y., Lin, D., & Mordasini, C. 2014, in Protostars and Planets VI, ed. H. Beuther, R. S. Klessen, C. P. Dullemond, & T. Henning, 691 Beuzit, J. L., Vigan, A., Mouillet, D., et al. 2019, arXiv e-prints, arXiv:1902.04080 Biller, B. A., Close, L. M., Masciadri, E., et al. 2007, ApJS, 173, 143
Birkby, J. L., de Kok, R. J., Brogi, M., et al. 2013, MNRAS, 436, L35
Bonaccini Calia, D., Feng, Y., Hackenberg, W., et al. 2010, The Messenger, 139, 12
Borucki, W. J., Koch, D., Basri, G., et al. 2010, Science, 327, 977 Boss, A. P. 1997, Science, 276, 1836
Bowler, B. P., Liu, M. C., Shkolnik, E. L., & Tamura, M. 2015, ApJS, 216, 7 Brogi, M., de Kok, R. J., Birkby, J. L., Schwarz, H., & Snellen, I. A. G. 2014,
A&A, 565, A124
Brogi, M., Snellen, I. A. G., de Kok, R. J., et al. 2013, ApJ, 767, 27 Burrows, A., Sudarsky, D., & Hubeny, I. 2004, ApJ, 609, 407
Butler, R. P., Marcy, G. W., Williams, E., Hauser, H., & Shirts, P. 1997, ApJ, 474, L115
Charbonneau, D., Noyes, R. W., Korzennik, S. G., et al. 1999, ApJ, 522, L145 Chauvin, G., Lagrange, A. M., Dumas, C., et al. 2005, A&A, 438, L25
BIBLIOGRAPHY 20
Close, L. M., Follette, K. B., Males, J. R., et al. 2014, ApJ, 781, L30
Codona, J. L., Kenworthy, M. A., Hinz, P. M., Angel, J. R. P., & Woolf, N. J. 2006, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 6269, 62691N
Crockett, C. J., Mahmud, N. I., Prato, L., et al. 2012, ApJ, 761, 164 Dodson-Robinson, S. E., & Salyk, C. 2011, ApJ, 738, 131
Foo, G., Palacios, D. M., & Swartzland er, Grover A., J. 2005, Optics Letters, 30, 3308
Foy, R., & Labeyrie, A. 1985, A&A, 152, L29
Fried, D. L. 1966, Journal of the Optical Society of America (1917-1983), 56, 1372 Fugate, R. Q., Fried, D. L., Ameer, G. A., et al. 1991, Nature, 353, 144
Gaudi, B. S., Stassun, K. G., Collins, K. A., et al. 2017, Nature, 546, 514 Gillon, M., Triaud, A. H. M. J., Demory, B.-O., et al. 2017, Nature, 542, 456 Greenwood, D. P. 1977, Journal of the Optical Society of America (1917-1983),
67, 390
Guyon, O. 2003, A&A, 404, 379
Haffert, S. Y. 2016, Optics Express, 24, 18986
Hennebelle, P., & Chabrier, G. 2008, ApJ, 684, 395, doi: 10.1086/589916 Henry, G. W., Marcy, G. W., Butler, R. P., & Vogt, S. S. 2000, ApJ, 529, L41 Hoeijmakers, H. J., Schwarz, H., Snellen, I. A. G., et al. 2018a, A&A, 617, A144 Hoeijmakers, H. J., Ehrenreich, D., Heng, K., et al. 2018b, Nature, 560, 453 Hubin, N., Arsenault, R., Conzelmann, R., et al. 2005, Comptes Rendus Physique,
6, 1099
Jovanovic, N., Martinache, F., Guyon, O., et al. 2015, PASP, 127, 890
Jovanovic, N., Absil, O., Baudoz, P., et al. 2018, in Society of Photo-Optical In-strumentation Engineers (SPIE) Conference Series, Vol. 10703, Adaptive Optics Systems VI, 107031U
Kasdin, N. J., Vanderbei, R. J., Spergel, D. N., & Littman, M. G. 2003, ApJ, 582, 1147
Kley, W. 2017, arXiv e-prints, arXiv:1707.07148
Konopacky, Q. M., Barman, T. S., Macintosh, B. A., & Marois, C. 2013, Science, 339, 1398
Lee, C.-H. 2017, Research Notes of the AAS, 1, 41
Lenzen, R., Hartung, M., Brandner, W., et al. 2003, in Proc. SPIE, Vol. 4841, Instrument Design and Performance for Optical/Infrared Ground-based Tele-scopes, ed. M. Iye & A. F. M. Moorwood, 944–952
Lovis, C., Snellen, I., Mouillet, D., et al. 2017, A&A, 599, A16 Lyot, B. 1939, MNRAS, 99, 580
Macintosh, B., Graham, J. R., Ingraham, P., et al. 2014, Proceedings of the Na-tional Academy of Science, 111, 12661
Madec, P. Y., Arsenault, R., Kuntschner, H., et al. 2018, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 10703, Adap-tive Optics Systems VI, 1070302
Marleau, G.-D., Klahr, H., Kuiper, R., & Mordasini, C. 2017, ApJ, 836, 221 Marois, C., Doyon, R., Nadeau, D., et al. 2005, PASP, 117, 745
Marois, C., Lafreni`ere, D., Doyon, R., Macintosh, B., & Nadeau, D. 2006, ApJ, 641, 556
Marois, C., Macintosh, B., Barman, T., et al. 2008, Science, 322, 1348
Marois, C., Zuckerman, B., Konopacky, Q. M., Macintosh, B., & Barman, T. 2010, Nature, 468, 1080
Martinez, P., Kasper, M., Costille, A., et al. 2013, A&A, 554, A41 Masuda, K. 2014, ApJ, 783, 53
Mayor, M., & Queloz, D. 1995, Nature, 378, 355 McKee, C. F., & Ostriker, E. C. 2007, ARA&A, 45, 565
Morbidelli, A., & Raymond, S. N. 2016, Journal of Geophysical Research (Planets), 121, 1962
Mulders, G. D., Pascucci, I., Apai, D., & Ciesla, F. J. 2018, AJ, 156, 24
Nelson, R. P., Papaloizou, J. C. B., Masset, F., & Kley, W. 2000, Monthly Notices of the Royal Astronomical Society, 318, 18
Nero, D., & Bjorkman, J. E. 2009, ApJ, 702, L163
Otten, G. P. P. L., Snik, F., Kenworthy, M. A., et al. 2017, ApJ, 834, 175 Pascucci, I., Testi, L., Herczeg, G. J., et al. 2016, ApJ, 831, 125
Petigura, E. A., Howard, A. W., & Marcy, G. W. 2013a, Proceedings of the National Academy of Science, 110, 19273
Petigura, E. A., Marcy, G. W., & Howard, A. W. 2013b, ApJ, 770, 69
Petit dit de la Roche, D. J. M., Hoeijmakers, H. J., & Snellen, I. A. G. 2018, A&A, 616, A146
Pollack, J. B., Hubickyj, O., Bodenheimer, P., et al. 1996, Icarus, 124, 62
Racine, R., Walker, G. A. H., Nadeau, D., Doyon, R., & Marois, C. 1999, PASP, 111, 587
Rameau, J., Chauvin, G., Lagrange, A. M., et al. 2013, A&A, 553, A60
Rasio, F. A., Shapiro, S. L., & Teukolsky, S. A. 1992, Astronomy and Astrophysics, 256, L35
Rouan, D., Riaud, P., Boccaletti, A., Cl´enet, Y., & Labeyrie, A. 2000, PASP, 112, 1479
Rousset, G., Lacombe, F., Puget, P., et al. 2003, in Proc. SPIE, Vol. 4839, Adaptive Optical System Technologies II, ed. P. L. Wizinowich & D. Bonaccini, 140–149 Ruane, G., Mawet, D., Jewell, J., & Shaklan, S. 2017, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 10400, Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, 104000J Ruane, G., Ngo, H., Mawet, D., et al. 2019, AJ, 157, 118
Schneider, G., & Silverstone, M. D. 2003, in Proc. SPIE, Vol. 4860, High-Contrast Imaging for Exo-Planet Detection., ed. A. B. Schultz, 1–9
Skemer, A. J., Marley, M. S., Hinz, P. M., et al. 2014, ApJ, 792, 17 Smith, B. A., & Terrile, R. J. 1984, Science, 226, 1421
Snellen, I., de Kok, R., Birkby, J. L., et al. 2015, A&A, 576, A59
BIBLIOGRAPHY 22
Snik, F., Otten, G., Kenworthy, M., et al. 2012, in Society of Photo-Optical Instru-mentation Engineers (SPIE) Conference Series, Vol. 8450, Modern Technologies in Space- and Ground-based Telescopes and Instrumentation II, 84500M Soummer, R. 2005, ApJ, 618, L161
Soummer, R., Aime, C., & Falloon, P. E. 2003, A&A, 397, 1161 Sparks, W. B., & Ford, H. C. 2002, ApJ, 578, 543
Tallon, M., & Foy, R. 1990, A&A, 235, 549
Tavani, M., & Brookshaw, L. 1992, Nature, 356, 320
Thatte, N., Abuter, R., Tecza, M., et al. 2007, MNRAS, 378, 1229
Traub, W. A., & Oppenheimer, B. R. 2010, Direct Imaging of Exoplanets (Uni-versity of Arizona Press), 111–156
van Eyken, J. C., Ciardi, D. R., von Braun, K., et al. 2012, ApJ, 755, 42 Vigan, A., Gry, C., Salter, G., et al. 2015, MNRAS, 454, 129
Vorobyov, E. I. 2013, A&A, 552, A129
Wilby, M. J., Keller, C. U., Haffert, S., et al. 2016, in Society of Photo-Optical In-strumentation Engineers (SPIE) Conference Series, Vol. 9909, Adaptive Optics Systems V, 990921
Wilby, M. J., Keller, C. U., Snik, F., Korkiakoski, V., & Pietrow, A. G. M. 2017, A&A, 597, A112
Wizinowich, P., Acton, D. S., Shelton, C., et al. 2000, PASP, 112, 315 Wolszczan, A., & Frail, D. A. 1992, Nature, 355, 145
